Giant Planet Accretion and Migration:

Surviving the Type I Regime

Giant Planet Accretion and Migration:

Surviving the Type I Regime

Edward Thommes

Norm Murray

CITA, University of Toronto

Edward Thommes

Norm Murray

CITA, University of Toronto

The Western Workshop, UWO, May 19, 2006The Western Workshop, UWO, May 19, 2006JPL

Gas giant formation: The core accretion model

Gas giant formation: The core accretion model

Gas disk lifetime sets upper limit on gas giant formation: ~1-10 Myrs from observations (e.g. Haisch, Lada & Lada 2001)

The core accretion model (Mizuno 1980, Pollack et al 1996):

1. Solid core grows, ~10 MEarth

2. Core accretes massive gas envelope, 100+ MEarth

Observational support for core accretion: planet-metallicity correlation

(Gonzalez 1997, Fischer & Valenti 2003)

HD 149026 planet (Saturn mass, ~70 MEarth core; Sato et al. 2005, Charbonneau et al 2006)

Gas disk lifetime sets upper limit on gas giant formation: ~1-10 Myrs from observations (e.g. Haisch, Lada & Lada 2001)

The core accretion model (Mizuno 1980, Pollack et al 1996):

1. Solid core grows, ~10 MEarth

2. Core accretes massive gas envelope, 100+ MEarth

Observational support for core accretion: planet-metallicity correlation

(Gonzalez 1997, Fischer & Valenti 2003)

HD 149026 planet (Saturn mass, ~70 MEarth core; Sato et al. 2005, Charbonneau et al 2006)

Marcy et al 2005

Planet-disk interactionPlanet-disk interaction Presence of substantial gas disk

means planet-disk interactions important!

Bodies in gas disk launch density waves repulsive torque between body and inner, outer disk

Jupiter-mass planets open a gap, locked into viscous evolution of disk: “Type II” inward migration

Smaller bodies: no gap, outer torques stronger: “Type I” inward migration

Presence of substantial gas disk means planet-disk interactions important!

Bodies in gas disk launch density waves repulsive torque between body and inner, outer disk

Jupiter-mass planets open a gap, locked into viscous evolution of disk: “Type II” inward migration

Smaller bodies: no gap, outer torques stronger: “Type I” inward migration

Density of planetdisk torque

Ward 1997

Geoff Bryden

Migration and accretion ratesMigration and accretion rates

Comparing the timescalesComparing the timescales

Scary result! Thus people tend to ignore/greatly reduce Type I (e.g. Thommes, Duncan & Levison 2003, Ida & Lin 2004, Alibert et al. 2005)

But is there a way to make the worst-case scenario work...?

Scary result! Thus people tend to ignore/greatly reduce Type I (e.g. Thommes, Duncan & Levison 2003, Ida & Lin 2004, Alibert et al. 2005)

But is there a way to make the worst-case scenario work...?

Accretion, no migration Accretion, no migration

Thommes & Murray 2006

Accretion + MigrationAccretion + Migration

Thommes & Murray 2006

A viscously evolving diskA viscously evolving disk

t=0

t=1 Myr

t=10 Myrs

Accretion + Migration in a viscously evolving gas diskAccretion + Migration in a viscously evolving gas disk

Thommes & Murray 2006

=10-2

Mdisk

M100 AU

M30 AU

Disk gas mass

Winners and losersWinners and losers Inner region:

growth too fast, cores lost onto star

Outer region: growth too slow relative to disk lifetime

In between: An annulus where the growth rate turns out just right

Inner region: growth too fast, cores lost onto star

Outer region: growth too slow relative to disk lifetime

In between: An annulus where the growth rate turns out just right

Thommes & Murray 2006

Method: Vary disk mass, metallicity,

For each set (MD,

,[Fe/H]), compute largest protoplanet mass when 1 MJup of gas left inside 100 AU

Results Disks with higher MD,

[Fe/H] do better There is always an

“optimal” , ~10-2-10-3; consistent with fits to T Tauri disks (Hartmann et al 1998)

Method: Vary disk mass, metallicity,

For each set (MD,

,[Fe/H]), compute largest protoplanet mass when 1 MJup of gas left inside 100 AU

Results Disks with higher MD,

[Fe/H] do better There is always an

“optimal” , ~10-2-10-3; consistent with fits to T Tauri disks (Hartmann et al 1998) Thommes & Murray 2006

Disk properties and core formationDisk properties and core formation

SummarySummary In the worst-case scenario of unmitigated Type I

migration: protoplanets in a young, massive gas disk fall onto

central star long before they can reach gas giant core size (~10 MEarth)...

...but as the gas disk dissipates, a window may open for cores to form and survive

endgame: gas envelope accretion plays large role in cleaning up rest of disk (cf. Lecar & Sasselov 2003)

Predictions Favourable disk properties: high M(0), high [Fe/H], and

~10-2 - 10-3

no giant planets (i.e. for ALMA, no gaps) in very young, massive disks

In the worst-case scenario of unmitigated Type I migration: protoplanets in a young, massive gas disk fall onto

central star long before they can reach gas giant core size (~10 MEarth)...

...but as the gas disk dissipates, a window may open for cores to form and survive

endgame: gas envelope accretion plays large role in cleaning up rest of disk (cf. Lecar & Sasselov 2003)

Predictions Favourable disk properties: high M(0), high [Fe/H], and

~10-2 - 10-3

no giant planets (i.e. for ALMA, no gaps) in very young, massive disks

“Dead zones” in disks“Dead zones” in disks Magnetorotational

instability (MRI) (Balbus & Hawley 1991) leading candidate for disk viscosity

MRI requires ionized disk, to couple it to magnetic field cosmic rays, stellar X-rays (near star)

When the full vertical column not ionized, dead zone forms (Gammie 1996, Matsumura & Pudritz 2003)

Magnetorotational instability (MRI) (Balbus & Hawley 1991) leading candidate for disk viscosity

MRI requires ionized disk, to couple it to magnetic field cosmic rays, stellar X-rays (near star)

When the full vertical column not ionized, dead zone forms (Gammie 1996, Matsumura & Pudritz 2003)

Gammie 1996

Disk evolution with a dead zoneDisk evolution with a dead zone

Dead zone: lower viscosityslower accretionpile-up of gas

Steep jumps in surface density can result

How does this affect migration...?

Dead zone: lower viscosityslower accretionpile-up of gas

Steep jumps in surface density can result

How does this affect migration...?

?

Disk torques at a surface density jump

Disk torques at a surface density jump

Type I migration: inner < outer, gas

Introducing jump in gas can reverse the torque imbalance

outer edge of a dead zone can completely stop Type I migration!

Type I migration: inner < outer, gas

Introducing jump in gas can reverse the torque imbalance

outer edge of a dead zone can completely stop Type I migration!

QuickTime™ and aTIFF (LZW) decompressor

are needed to see this picture.

Matsumura, Thommes & Pudritz, in prep.

QuickTime™ and aCinepak decompressor

are needed to see this picture.

Thommes

A “hybrid” code: N-body+gas diskA “hybrid” code: N-body+gas disk The N-body part: SyMBA (Duncan, Levison & Lee 1998)

uses Wisdom-Holman (1991) symplectic method fast for near-Keplerian systems bounded energy error

resolves close encounters The disk-evolution part: 1-D (azimuthally, vertically

averaged) Keplerian disk, Σ evolves according to

The N-body part: SyMBA (Duncan, Levison & Lee 1998) uses Wisdom-Holman (1991) symplectic method

fast for near-Keplerian systems bounded energy error

resolves close encounters The disk-evolution part: 1-D (azimuthally, vertically

averaged) Keplerian disk, Σ evolves according to

(Goldreich & Tremaine 1980, Ward 1997)

--∫(dT/dr)dr applied to planet∫(dT/dr)dr applied to planet ...Fast! Can simulate 10...Fast! Can simulate 1077 yrs in ~2 days yrs in ~2 days

-

QuickTime™ and aCinepak decompressor

are needed to see this picture.

Thommes 2005

Resonant exoplanetsResonant exoplanets

Marcy et al. 2005

The “standard model” of core accretion

The “standard model” of core accretion

Pollack et al (1996): 3 stages:1. solid core accretion

2. slow gas accretion until Mgas ~ Mcore

3. runaway gas accretion Corrections to the standard

model: Stage 1 simplified, actually

takes longer (e.g. Thommes et al. 2003)

Stage 2 HAS to be a lot shorter (can be done by lowering envelope opacity)

Pollack et al (1996): 3 stages:1. solid core accretion

2. slow gas accretion until Mgas ~ Mcore

3. runaway gas accretion Corrections to the standard

model: Stage 1 simplified, actually

takes longer (e.g. Thommes et al. 2003)

Stage 2 HAS to be a lot shorter (can be done by lowering envelope opacity)

Pollack et al. 1996

OutlineOutline Background

giant planet formation by core accretion migration by planet-disk interaction

The timescale problem Calculations of concurrent core accretion and migration in

an evolving disk A way around the timescale problem

Disk properties and the prospects for planet formation Summary

Background giant planet formation by core accretion migration by planet-disk interaction

The timescale problem Calculations of concurrent core accretion and migration in

an evolving disk A way around the timescale problem

Disk properties and the prospects for planet formation Summary